Venus has thick clouds of H2SO4 aerosol particles
extending from altitudes of 40 to 60km. The 60-100km region (the
mesosphere) is a transition region between the 4day retrograde
superrotation at the top of the thick clouds and the solar-antisolar
circulation in the thermosphere (above 100km), which has upwelling over
the subsolar point and transport to the nightside. The mesosphere has a
light haze of variable optical thickness, with CO, SO2, HCl,
HF, H2O and HDO as the most important minor gaseous
constituents, but the vertical distribution of the haze and molecules is
poorly known because previous descent probes began their measurements at
or below 60km. Here we report the detection of an extensive layer of
warm air at altitudes 90-120km on the night side that we interpret as
the result of adiabatic heating during air subsidence. Such a strong
temperature inversion was not expected, because the night side of Venus
was otherwise so cold that it was named the `cryosphere' above 100km. We
also measured the mesospheric distributions of HF, HCl, H2O
and HDO. HCl is less abundant than reported 40years ago.
HDO/H2O is enhanced by a factor of ˜2.5 with respect to the
lower atmosphere, and there is a general depletion of H2O
around 80-90km for which we have no explanation.

The upper atmosphere of a planet is a transition region in which energy
is transferred between the deeper atmosphere and outer space. Molecular
emissions from the upper atmosphere (90-120km altitude) of Venus can be
used to investigate the energetics and to trace the circulation of this
hitherto little-studied region. Previous spacecraft and ground-based
observations of infrared emission from CO2, O2 and
NO have established that photochemical and dynamic activity controls the
structure of the upper atmosphere of Venus. These data, however, have
left unresolved the precise altitude of the emission owing to a lack of
data and of an adequate observing geometry. Here we report measurements
of day-side CO2 non-local thermodynamic equilibrium emission
at 4.3m, extending from 90 to 120km altitude, and of night-side
O2 emission extending from 95 to 100km. The CO2
emission peak occurs at ˜115km and varies with solar zenith angle over a
range of ˜10km. This confirms previous modelling, and permits the
beginning of a systematic study of the variability of the emission. The
O2 peak emission happens at 96km+/-1km, which is consistent
with three-body recombination of oxygen atoms transported from the day
side by a global thermospheric sub-solar to anti-solar circulation, as
previously predicted.

Venus has no seasons, slow rotation and a very massive atmosphere, which
is mainly carbon dioxide with clouds primarily of sulphuric acid
droplets. Infrared observations by previous missions to Venus revealed a
bright `dipole' feature surrounded by a cold `collar' at its north pole.
The polar dipole is a `double-eye' feature at the centre of a vast
vortex that rotates around the pole, and is possibly associated with
rapid downwelling. The polar cold collar is a wide, shallow river of
cold air that circulates around the polar vortex. One outstanding
question has been whether the global circulation was symmetric, such
that a dipole feature existed at the south pole. Here we report
observations of Venus' south-polar region, where we have seen clouds
with morphology much like those around the north pole, but rotating
somewhat faster than the northern dipole. The vortex may extend down to
the lower cloud layers that lie at about 50km height and perhaps deeper.
The spectroscopic properties of the clouds around the south pole are
compatible with a sulphuric acid composition.

A unique feature of the Martian climate is the possibility for carbon
dioxide, the main atmospheric constituent, to condense as ice.
CO2 ice is usually detected as frost but is also known to
exist as clouds. This paper presents the first unambiguous observation
of CO2 ice clouds on Mars. These images were obtained by the
visible and near-infrared imaging spectrometer OMEGA on board Mars
Express. The data set encompasses 19 different occurrences.
Compositional identification is based on the detection of a diagnostic
spectral feature around 4.26 μm which is produced by resonant
scattering of solar photons by mesospheric CO2 ice particles
in a spectral interval otherwise dominated by saturated gaseous
absorption. Observed clouds exhibit a strong seasonal and geographic
dependence, concentrating in the near-equatorial regions during two
periods before and after northern summer solstice (Ls 45deg and
135deg). Radiative transfer modeling indicates that the 4.26 μm
feature is very sensitive to cloud altitude, opacity, and particle size,
thereby explaining the variety of spectra associated with the cloud
images. On two orbits, the simultaneous detection of clouds with their
shadow provides straightforward and robust estimates of cloud
properties. These images confirm the conclusions established from
modeling: clouds are thick, with normal opacities greater than 0.2 in
the near infrared, and are lofted in the mesosphere above 80 km. The
mean radius of CO2 ice crystals is found to exceed 1 μm,
an unexpected value considering this altitude range. This finding
implies the existence of high-altitude atmospheric updrafts which are
strong enough to counteract the rapid gravitational fall of particles.
This statement is consistent with the cumuliform morphology of the
clouds which may be linked to a moist convective origin generated by the
latent heat released during CO2 condensation.

Spectroscopy for the investigation of the characteristics of the
atmosphere of Venus (SPICAV) is a suite of three spectrometers in the UV
and IR range with a total mass of 13.9 kg flying on the Venus Express
(VEX) orbiter, dedicated to the study of the atmosphere of Venus from
ground level to the outermost hydrogen corona at more than 40,000 km. It
is derived from the SPICAM instrument already flying on board Mars
Express (MEX) with great success, with the addition of a new IR
high-resolution spectrometer, solar occultation IR (SOIR), working in
the solar occultation mode. The instrument consists of three
spectrometers and a simple data processing unit providing the interface
of these channels with the spacecraft. A UV spectrometer (118-320 nm,
resolution 1.5 nm) is identical to the MEX version. It is dedicated to
nadir viewing, limb viewing and vertical profiling by stellar and solar
occultation. In nadir orientation, SPICAV UV will analyse the albedo
spectrum (solar light scattered back from the clouds) to retrieve SO
2, and the distribution of the UV-blue absorber (of still
unknown origin) on the dayside with implications for cloud structure and
atmospheric dynamics. On the nightside, γ and δ bands of NO
will be studied, as well as emissions produced by electron
precipitations. In the stellar occultation mode the UV sensor will
measure the vertical profiles of CO 2, temperature, SO
2, SO, clouds and aerosols. The density/temperature profiles
obtained with SPICAV will constrain and aid in the development of
dynamical atmospheric models, from cloud top (60 km) to 160 km in
the atmosphere. This is essential for future missions that would rely on
aerocapture and aerobraking. UV observations of the upper atmosphere
will allow studies of the ionosphere through the emissions of CO, CO
+, and CO 2+, and its direct
interaction with the solar wind. It will study the H corona, with its
two different scale heights, and it will allow a better understanding of
escape mechanisms and estimates of their magnitude, crucial for insight
into the long-term evolution of the atmosphere. The SPICAV VIS-IR sensor
(0.7-1.7 μm, resolution 0.5-1.2 nm) employs a pioneering technology:
an acousto-optical tunable filter (AOTF). On the nightside, it will
study the thermal emission peeping through the clouds, complementing the
observations of both VIRTIS and Planetary Fourier Spectrometer (PFS) on
VEX. In solar occultation mode this channel will study the vertical
structure of H 2O, CO 2, and aerosols. The SOIR
spectrometer is a new solar occultation IR spectrometer in the range
λ=2.2-4.3 μm, with a spectral resolution λ/Δ
λ15,000, the highest on board VEX. This new concept includes
a combination of an echelle grating and an AOTF crystal to sort out one
order at a time. The main objective is to measure HDO and H
2O in solar occultation, in order to characterize the escape
of D atoms from the upper atmosphere and give more insight about the
evolution of water on Venus. It will also study isotopes of CO
2 and minor species, and provides a sensitive search for new
species in the upper atmosphere of Venus. It will attempt to measure
also the nightside emission, which would allow a sensitive measurement
of HDO in the lower atmosphere, to be compared to the ratio in the upper
atmosphere, and possibly discover new minor atmospheric constituents.

The Visible and Infrared Thermal Imaging Spectrometer (VIRTIS) on board
the ESA/Venus Express mission has technical specifications well suited
for many science objectives of Venus exploration. VIRTIS will both
comprehensively explore a plethora of atmospheric properties and
processes and map optical properties of the surface through its three
channels, VIRTIS-M-vis (imaging spectrometer in the 0.3-1 μm range),
VIRTIS-M-IR (imaging spectrometer in the 1-5 μm range) and VIRTIS-H
(aperture high-resolution spectrometer in the 2-5 μm range). The
atmospheric composition below the clouds will be repeatedly measured in
the night side infrared windows over a wide range of latitudes and
longitudes, thereby providing information on Venus's chemical cycles. In
particular, CO, H 2O, OCS and SO 2 can be studied.
The cloud structure will be repeatedly mapped from the brightness
contrasts in the near-infrared night side windows, providing new
insights into Venusian meteorology. The global circulation and local
dynamics of Venus will be extensively studied from infrared and visible
spectral images. The thermal structure above the clouds will be
retrieved in the night side using the 4.3 μm fundamental band of CO
2. The surface of Venus is detectable in the short-wave
infrared windows on the night side at 1.01, 1.10 and 1.18 μm,
providing constraints on surface properties and the extent of active
volcanism. Many more tentative studies are also possible, such as
lightning detection, the composition of volcanic emissions, and
mesospheric wave propagation.

We present the seasonal and geographical variations of the martian water
vapor monitored from the Planetary Fourier Spectrometer Long Wavelength
Channel aboard the Mars Express spacecraft. Our dataset covers one
martian year (end of Mars Year 26, Mars Year 27), but the seasonal
coverage is far from complete. The seasonal and latitudinal behavior of
the water vapor is globally consistent with previous datasets, Viking
Orbiter Mars Atmospheric Water Detectors (MAWD) and Mars Global Surveyor
Thermal Emission Spectrometer (MGS/TES), and with simultaneous results
obtained from other Mars Express instruments, OMEGA and SPICAM. However,
our absolute water columns are lower and higher by a factor of 1.5 than
the values obtained by TES and SPICAM, respectively. In particular, we
retrieve a Northern midsummer maximum of 60 pr-μm, lower than the
100-pr-μm observed by TES. The geographical distribution of water
exhibits two local maxima at low latitudes, located over Tharsis and
Arabia. Global Climate Model (GCM) simulations suggest that these local
enhancements are controlled by atmospheric dynamics. During Northern
spring, we observe a bulge of water vapor over the seasonal polar cap
edge, consistent with the northward transport of water from the
retreating seasonal cap to the permanent polar cap. In terms of vertical
distribution, we find that the water volume mixing ratio over the large
volcanos remains constant with the surface altitude within a factor of
two. However, on the whole dataset we find that the water column,
normalized to a fixed pressure, is anti-correlated with the surface
pressure, indicating a vertical distribution intermediate between
control by atmospheric saturation and confinement to a surface layer.
This anti-correlation is not reproduced by GCM simulations of the water
cycle, which do not include exchange between atmospheric and subsurface
water. This situation suggests a possible role for regolith-atmosphere
exchange in the martian water cycle.

The poles of Mars are known to have recorded recent (107
years) climatic changes. While the south polar region appears to have
preserved its million-year-old environment from major resurfacing
events, except for the small portion containing the CO2
residual cap, the discovery of residual water ice units in areas
adjacent to the cap provides compelling evidence for recent
glaciological activity. The mapping and characterization of these
H2O-rich terrains by Observatoire pour la Minéralogie,
l'Eau, les Glaces et l'Activité (OMEGA) on board Mars Express,
which have supplemented earlier findings by Mars Odyssey and Mars Global
Surveyor, have raised a number of questions related to their origin. We
propose that these water ice deposits are the relics of Mars' orbit
precession cycle and that they were laid down when perihelion was
synchronized with northern summer, i.e., more than 10,000 years ago. We
favor precession over other possible explanations because (1) as shown
by our General Circulation Model (GCM) and previous studies, current
climate is not conducive to the accumulation of water at the south pole
due to an unfavorable volatile transport and insolation configuration,
(2) the residual CO2 ice cap, which is known to cold trap
water molecules on its surface and which probably controls the current
extent of the water ice units, is geologically younger, (3) our GCM
shows that 21,500 years ago, when perihelion occurred during northern
spring, water ice at the north pole was no longer stable and accumulated
instead near the south pole with rates as high as 1 mm yr-1.
This could have led to the formation of a meters-thick circumpolar water
ice mantle. As perihelion slowly shifted back to the current value,
southern summer insolation intensified and the water ice layer became
unstable. The layer recessed poleward until the residual CO2
ice cover eventually formed on top of it and protected water ice from
further sublimation. In this polar accumulation process, water ice
clouds play a critical role since they regulate the exchange of water
between hemispheres. The so-called “Clancy effect,” which sequesters
water in the spring/summer hemisphere coinciding with aphelion due to
cloud sedimentation, is demonstrated to be comparable in magnitude to
the circulation bias forced by the north-to-south topographic dichotomy.
However, we predict that the response of Mars' water cycle to the
precession cycle should be asymmetric between hemispheres not only
because of the topographic bias in circulation but also because of an
asymmetry in the dust cycle. We predict that under a “reversed
perihelion” climate, dust activity during northern summer is less
pronounced than during southern summer in the opposite perihelion
configuration (i.e., today's regime). When averaged over a precession
cycle, this reduced potential for dust lifting will force a
significantly colder summer in the north and, by virtue of the Clancy
effect, will curtail the ability of the northern hemisphere to transfer
volatiles to the south. This process may have helped create the observed
morphological differences in the layered deposits between the poles and
could help explain the large disparity in their resurfacing ages.

Surface pressure measurements help to achieve a better understanding of
the main dynamical phenomena occurring in the atmosphere of a planet.
The use of the Mars Express OMEGA visible and near-IR imaging
spectrometer allows us to tentatively perform an unprecedented remote
sensing measurement of Martian surface pressure. OMEGA reflectances in
the CO2 absorption band at 2 μm are used to retrieve a
hydrostatic estimation of surface pressure (see companion paper by
Forget et al. (2007)) with a precision sufficient to draw maps of this
field and thus analyze meteorological events in the Martian atmosphere.
Prior to any meteorological analysis, OMEGA observations have to pass
quality controls on insolation and albedo conditions, atmosphere dust
opacity, and occurrence of water ice clouds and frosts. For the selected
observations, registration shifts with the MOLA reference are corrected.
“Sea-level” surface pressure reduction is then carried out in order to
remove the topographical component of the surface pressure field. Three
main phenomena are observed in the resulting OMEGA surface pressure
maps: horizontal pressure gradients, atmospheric oscillations, and
pressure perturbations in the vicinity of topographical obstacles. The
observed pressure oscillations are identified as possible signatures of
phenomena such as inertia-gravity waves or convective rolls. The
pressure perturbations detected around the Martian hills and craters may
be the signatures of complex interactions between an incoming flow and
topographical obstacles. Highly idealized mesoscale simulations using
the WRF model enable a preliminary study of these complex interactions,
but more realistic mesoscale simulations are necessary. The maps provide
valuable insights for future synoptic and mesoscale modeling, which will
in turn help in the interpretation of observations.

Observing and analyzing the variations of pressure on the surface of a
planet is essential to understand the dynamics of its atmosphere. On
Mars the absorption by atmospheric CO2 of the solar light
reflected on the surface allows us to measure the surface pressure by
remote sensing. We use the imaging spectrometer OMEGA aboard Mars
Express, which provides an excellent signal to noise ratio and the
ability to produce maps of surface pressure with a resolution ranging
from 400 m to a few kilometers. Surface pressure is measured by fitting
spectra of the CO2 absorption band centered at 2 μm. To
process the hundreds of thousands of pixels present in each OMEGA image,
we have developed a fast and accurate algorithm based on a line-by-line
radiative transfer model which includes scattering and absorption by
dust aerosols. In each pixel the temperature profile, the dust opacity,
and the surface spectrum are carefully determined from the OMEGA data
set or from other sources to maximize the accuracy of the retrieval. We
estimate the 1-σ relative error to be around 7 Pa in bright
regions and about 10 Pa in darker regions, with a possible systematic
bias on the absolute pressure lower than 30 Pa (4%). The method is first
tested by comparing an OMEGA pressure retrieval obtained over the Viking
Lander 1 (VL1) landing site with in situ measurements recorded 30 years
ago by the VL1 barometer. The retrievals are further validated using a
surface pressure predictor which combines the VL1 pressure records with
the MOLA topography and meteorological pressure gradients simulated with
a General Circulation Model. A good agreement is obtained. In
particular, OMEGA is able to monitor the seasonal variations of the
surface pressure in Isidis Planitia. Such a tool can be applied to
detect meteorological phenomena, as described by Spiga et al. (2007).

The OMEGA visible/near-infrared imaging spectrometer on board Mars
Express has observed the southern seasonal cap in late 2004 and 2005 and
then in the summer of 2006. These observations extended from the period
of maximum extension, close to the southern winter solstice, to the end
of the recession at Ls 325deg. The spectral range and
spectral resolution of OMEGA make it possible to monitor the extent and
effective grain size of CO2 ice and H2O ice on the
ground, the level of contamination of CO2 ice and
H2O ice by dust, and the column density of μm-sized ice
grains in the atmosphere. The CO2 seasonal cap is very clean
and clear in early southern winter. Contamination by H2O ice
spreads eastward from the Hellas basin until the southern spring
equinox. During southern spring and summer, there is a very complex
evolution in terms of effective grain size of CO2 ice and
contamination by dust or H2O ice. H2O ice does not
play a significant role close to the southern summer solstice.
Contamination of CO2 ice by H2O ice is only
observed close to the end of the recession, as well as the few
H2O ice patches already reported by Bibring et al. (2004a).
These observations have been compared to the results of a general
circulation model, with good qualitative agreement on the distribution
of H2O ice on the surface and in the atmosphere. Resolving
the remaining discrepancies will improve our understanding of the water
cycle on Mars.

This work presents a review of the observations acquired by the
planetary Fourier spectrometer (PFS) in the region of the Hellas basin.
Taking advantage of the high spectral resolution of PFS, the vertical
air temperature profile can be investigated with a previously
unexperienced vertical resolution. Extensive comparisons with the
expectations of EMCD 4.0 database highlight moderate discrepancies,
strongly dependant on season. Namely, the morning observations acquired
around Ls=45deg show a series of temperature deficiencies
with recurrent spatial patterns in different observations, correlated
with the topography profile. Trends of integrated dust loads as a
function of the field of view (FOV) elevation are also described. Values
are consistent with the retrieval hypothesis of a dust scale height
equal to the gas one, even far from the season of main dust storms.

We present a time-marching model which simulates the exchange of water
ice between the Martian northern cap, the tropics, and a high-latitude
surface reservoir. Net annual exchange rates of water and their
sensitivity to variations in orbital/rotational parameters are examined
using the Martian water cycle modeled by the LMD three-dimensional
Global Climate Model. These rates are propagated over the last 10 Myr to
follow the thickness of the reservoirs. The effect of a sublimation dust
lag is taken account to test simple models of layer formation. Periods
of high mean polar summer insolation (˜5-10 Ma ago) lead to a rapid
exhaustion of a northern polar cap and a prolonged formation of tropical
glaciers. The formation of a northern cap and of a high-latitude icy
mantle may have started 4 Ma ago with the average decrease of polar
insolation. Tropical ice may have disappeared around 2.7 Ma ago, but
small glaciers could have formed during the last peaks of polar summer
insolation. Over the last 4 Myr, most of the present cap may have formed
at the expense of tropical and high-latitude reservoirs forming distinct
layers at almost each ˜51-kyr/120-kyr insolation cycle. Layers thickness
ranges from 10 to 80 m, variations being produced by the modulation of
the obliquity with ˜2.4 and 1.3 Myr periods. Because only ˜30 insolation
cycles have occurred since 4 Ma ago, we found an inconsistency between
the recent astronomical forcing, the observed number of layers, and
simple astronomically based scenarios of layers formation.